Oxygen-burning process

The oxygen-burning process is a set of nuclear fusion reactions that take place in massive stars that have used up the lighter elements in their cores. Oxygen-burning is preceded by the neon-burning process and succeeded by the silicon-burning process. As the neon-burning process ends, the core of the star contracts and heats until it reaches the ignition temperature for oxygen burning. Oxygen burning reactions are similar to those of carbon burning; however, they must occur at higher temperatures and densities due to the larger Coulomb barrier of oxygen. Oxygen in the core ignites in the temperature range of (1.5–2.6)×109 K[1] and in the density range of (2.6–6.7)×109g/cm3.[2] The principal reactions are given below,[3][4] where the branching ratios assume that the deuteron channel is open (at high temperatures):[3]

→  28
9.593 MeV (34%)
      →  31
7.676 MeV (56%)
      →  31

1.459 MeV (5%)
      →  30
0.381 MeV
      →  30
2.409 MeV (5%)
      →  32

+ 16.539 MeV
      →  24
- 0.393 MeV


Near 2×109K, the oxygen burning reaction rate is approximately 2.8×10−12(T9/2)33,[3][5] where T9 is the temperature in billions of Kelvin. Overall, the major products of the oxygen-burning process are [3] 28Si, 32,33,34S, 35,37Cl, 36,38Ar, 39,41K, and 40,42Ca. Of these, 28Si and 32S constitute 90% of the final composition.[3] The oxygen fuel within the core of the star is exhausted after 0.01–5 years depending on the star's mass and other parameters.[1][3] The silicon-burning process which follows creates iron, but this iron cannot react further to create energy to support the star.

During the oxygen-burning process, proceeding outward, there is an oxygen-burning shell, followed by a neon shell, a carbon shell, a helium shell, and a hydrogen shell. The oxygen-burning process is the last nuclear reaction in the star's core which does not proceed via the alpha process.

Pre-oxygen burning

Although 16O is lighter than neon, neon burning occurs before oxygen burning, because 16O is a doubly-magic nucleus and hence extremely stable. Compared to oxygen, neon is much less stable. As a result, neon burning occurs at lower temperatures than 16O+16O.[9] During neon burning, oxygen and magnesium accumulate in the core of the star. At the onset of oxygen burning, oxygen in the stellar core is plentiful due to the helium-burning process (4He(2α,γ)12C(α,γ)16O), carbon-burning process (12C(12C,α)20Ne, 12C(α,γ)16O), and neon-burning process (20Ne(γ,α)16O). The reaction 12C(α,γ)16O has a significant effect on the reaction rates during oxygen burning, as it produces large quantities of 16O.[3]

Convectively bounded flames and off-center oxygen ignition

For stars with masses greater than 10.3 solar masses, oxygen ignites in the core or not at all. Similarly, for stars with a mass of less than 9 solar masses (without accretion of additional mass) oxygen ignites in the core or not at all. However, in the 9–10.3 solar mass range, oxygen ignites off-center.

For stars in this mass range neon-burning occurs in a convective envelope rather than at the core of the star. For the particular example of a 9.5 solar mass star, the neon-burning process takes place in an envelope of approximately 0.252 solar masses (~1560 kilometers) off center. From the ignition flash, the neon convective zone extends further out to 1.1 solar masses with a peak power around 1043 erg/s. After only a month, the power declines to about 1042 erg/s and stays at this rate for about 10 years. After this phase, the neon in the shell is depleted, resulting in greater inward pressure on the star. This raises the shell's temperature to 1.65 billion Kelvin. This results in a neon-burning, convectively-bound flame front that moves toward the core. The motion of the flame is what eventually leads to oxygen-burning. In approximately 3 years, the flame's temperature reaches about 1.83 billion Kelvin, enabling the oxygen-burning process to commence. This occurs around 9.5 years before the iron core develops. Similarly to the beginning of neon-burning, off-center oxygen-burning commences with another flash. The convectively-burning flame then results from both neon and oxygen burning as it advances towards the core, while the oxygen-burning shell continuously shrinks in mass.[8]

Neutrino losses

During the oxygen-burning process, energy loss due to neutrino emission becomes relevant. Due to the large energy loss, oxygen must burn at temperatures higher than a billion Kelvin in order to maintain a radiation pressure strong enough to support the star against gravity. Further, two electron capture reactions (which produce neutrinos) become significant when the matter density is high enough (ρ>2×107 g•cm−3). Due to these factors, the timescale of oxygen burning is much shorter for heavy, dense stars.[7]

Explosive oxygen burning

The oxygen-burning process can occur under hydrostatic and under explosive conditions. The products of explosive oxygen burning are similar to those in hydrostatic oxygen burning. However, stable oxygen burning is accompanied by a multitude of electron captures, while explosive oxygen burning is accompanied by a significantly greater presence of photodisintegration reactions. In the temperature range of (3–4)×109 K, photodisintegration and oxygen fusion occur with comparable reaction rates.[3]

Pair-instability supernovae

Very massive (140–260 solar masses) population III stars may become unstable during core oxygen burning due to pair production. This results in a thermonuclear explosion which completely disrupts the star.[2][6]


  1. ^ a b Eid, M. F. El, B. S. Meyer, and L.‐S. The. "Evolution of Massive Stars Up to the End of Central Oxygen Burning." ApJ The Astrophysical Journal 611.1 (2004): 452–65. Arxiv.org. 21 July 2004. Web. 8 Apr. 2016.
  2. ^ a b Hirschi. “Evolution and nucleosynthesis of Very Massive Stars.” arXiv:1409.7053v1 [astro-ph.SR] 24 Sep 2014
  3. ^ a b c d e f g h Woosley, Heger, and Weaver. “The evolution of massive stars.” Reviews of Modern Physics, Volume 74, October 2002
  4. ^ Clayton, Donald. Principles of Stellar Evolution and Nucleosynthesis, (1983)
  5. ^ a b Caughlan and Fowler. "Thermonuclear reaction rates”. Atomic Data and Nuclear Data Tables, 40, 283–334 (1988).
  6. ^ a b Kasen, Woosley, and Heger. “Pair Instability Supernovae: Light Curves, Spectra, and Shock Breakout.” The Astrophysical Journal 734:102, 2011 June 20.
  7. ^ a b Carroll, Bradley W., and Dale A. Ostlie." An Introduction to Modern Astrophysics." San Francisco, Pearson Addison-Wesley, 2007. Print.
  8. ^ a b S.E. Woosley and Alexander Heger. “The Remarkable Deaths of 9–10 Solar Mass Stars.” arXiv:1505.06712v1. May 2015.
  9. ^ a b Longair, Malcolm. High Energy Astrophysics, 3rd Edition, (2011)

External links

B2FH paper

The B2FH paper, named after the initials of the authors of the paper, Margaret Burbidge, Geoffrey Burbidge, William A. Fowler, and Fred Hoyle, is a landmark paper on the origin of the chemical elements published in Reviews of Modern Physics in 1957. The title of that paper is "Synthesis of the Elements in Stars", but as that paper grew in influence, it came to be referred to only as "B2FH". The B2FH paper spread stellar nucleosynthesis theory widely in the scientific community, especially among astronomers who saw everyday relevance to their quest, at a time when it was appreciated by only a handful of experts in nuclear physics. But it did not create the theory of stellar nucleosynthesis as much as bring it vividly to life.

The paper comprehensively outlined and analyzed several key processes that are responsible for the nucleosynthesis of the elements heavier than iron and their relative abundance by the capture within stars of free neutrons. It advanced much less the understanding of the synthesis of the very abundant elements from silicon to nickel. A puzzle about that is that despite Hoyle's coauthorship of B2FH and being its chief conceptual architect, the paper did not include the carbon-burning process, the oxygen-burning process and the silicon-burning process, each of which contributes massively to the growth of stellar metallicity from magnesium to nickel in the interstellar gas. The supernova nucleosynthesis that achieves that had been published by Hoyle in 1954. Donald D. Clayton has attributed the severe undercitations of Hoyle's 1954 paper relative to the voluminous citations of B2FH to several factors: the advanced difficulty of digesting Hoyle's 1954 paper even for his B2FH coauthors, as it proved to be for the world of astronomy generally; to Hoyle's having described its key equation only in words rather than writing it prominently in his paper; and finally to a lack of careful review by Hoyle himself of the B2FH draft written by two junior coauthors who had themselves not adequately digested Hoyle's paper.

Index of physics articles (O)

The index of physics articles is split into multiple pages due to its size.

To navigate by individual letter use the table of contents below.

Neon-burning process

The neon-burning process (nuclear decay) is a set of nuclear fusion reactions that take place in massive stars (at least 8 Solar masses). Neon burning requires high temperatures and densities (around 1.2×109 K or 100 KeV and 4×109 kg/m3).

At such high temperatures photodisintegration becomes a significant effect, so some neon nuclei decompose, releasing alpha particles:


where the neutron consumed in the first step is regenerated in the second.

Neon burning takes place after carbon burning has consumed all carbon in the core and built up a new oxygen-neon-sodium-magnesium core. The core ceases producing fusion energy and contracts. This contraction increases density and temperature up to the ignition point of neon burning. The increased temperature around the core allows carbon to burn in a shell, and there will be shells burning helium and hydrogen outside.

During neon burning, oxygen and magnesium accumulate in the central core while neon is consumed. After a few years the star consumes all its neon and the core ceases producing fusion energy and contracts. Again, gravitational pressure takes over and compresses the central core, increasing its density and temperature until the oxygen-burning process can start.

Outline of astronomy

The following outline is provided as an overview of and topical guide to astronomy:

Astronomy – studies the universe beyond Earth, including its formation and development, and the evolution, physics, chemistry, meteorology, and motion of celestial objects (such as galaxies, planets, etc.) and phenomena that originate outside the atmosphere of Earth (such as the cosmic background radiation).


Potassium is a chemical element with the symbol K (from Neo-Latin kalium) and atomic number 19. Potassium is a silvery-white metal that is soft enough to be cut with a knife with little force. Potassium metal reacts rapidly with atmospheric oxygen to form flaky white potassium peroxide in only seconds of exposure. It was first isolated from potash, the ashes of plants, from which its name derives. In the periodic table, potassium is one of the alkali metals, all of which have a single valence electron in the outer electron shell, that is easily removed to create an ion with a positive charge – a cation, that combines with anions to form salts. Potassium in nature occurs only in ionic salts. Elemental potassium reacts vigorously with water, generating sufficient heat to ignite hydrogen emitted in the reaction, and burning with a lilac-colored flame. It is found dissolved in sea water (which is 0.04% potassium by weight), and occurs in many minerals such as orthoclase, a common constituent of granites and other igneous rocks.

Potassium is chemically very similar to sodium, the previous element in group 1 of the periodic table. They have a similar first ionization energy, which allows for each atom to give up its sole outer electron. That they are different elements that combine with the same anions to make similar salts was suspected in 1702, and was proven in 1807 using electrolysis. Naturally occurring potassium is composed of three isotopes, of which 40K is radioactive. Traces of 40K are found in all potassium, and it is the most common radioisotope in the human body.

Potassium ions are vital for the functioning of all living cells. The transfer of potassium ions across nerve cell membranes is necessary for normal nerve transmission; potassium deficiency and excess can each result in numerous signs and symptoms, including an abnormal heart rhythm and various electrocardiographic abnormalities. Fresh fruits and vegetables are good dietary sources of potassium. The body responds to the influx of dietary potassium, which raises serum potassium levels, with a shift of potassium from outside to inside cells and an increase in potassium excretion by the kidneys.

Most industrial applications of potassium exploit the high solubility in water of potassium compounds, such as potassium soaps. Heavy crop production rapidly depletes the soil of potassium, and this can be remedied with agricultural fertilizers containing potassium, accounting for 95% of global potassium chemical production.


Silicon is a chemical element with the symbol Si and atomic number 14. It is a hard and brittle crystalline solid with a blue-grey metallic lustre; and it is a tetravalent metalloid and semiconductor. It is a member of group 14 in the periodic table: carbon is above it; and germanium, tin, and lead are below it. It is relatively unreactive. Because of its high chemical affinity for oxygen, it was not until 1823 that Jöns Jakob Berzelius was first able to prepare it and characterize it in pure form. Its melting and boiling points of 1414 °C and 3265 °C respectively are the second-highest among all the metalloids and nonmetals, being only surpassed by boron. Silicon is the eighth most common element in the universe by mass, but very rarely occurs as the pure element in the Earth's crust. It is most widely distributed in dusts, sands, planetoids, and planets as various forms of silicon dioxide (silica) or silicates. More than 90% of the Earth's crust is composed of silicate minerals, making silicon the second most abundant element in the Earth's crust (about 28% by mass) after oxygen.

Most silicon is used commercially without being separated, and often with little processing of the natural minerals. Such use includes industrial construction with clays, silica sand, and stone. Silicates are used in Portland cement for mortar and stucco, and mixed with silica sand and gravel to make concrete for walkways, foundations, and roads. They are also used in whiteware ceramics such as porcelain, and in traditional quartz-based soda-lime glass and many other specialty glasses. Silicon compounds such as silicon carbide are used as abrasives and components of high-strength ceramics. Silicon is the basis of the widely used synthetic polymers called silicones.

Elemental silicon also has a large impact on the modern world economy. Most free silicon is used in the steel refining, aluminium-casting, and fine chemical industries (often to make fumed silica). Even more visibly, the relatively small portion of very highly purified elemental silicon used in semiconductor electronics (< 10%) is essential to integrated circuits – most computers, cell phones, and modern technology depend on it.

Silicon is an essential element in biology, although only traces are required by animals. However, various sea sponges and microorganisms, such as diatoms and radiolaria, secrete skeletal structures made of silica. Silica is deposited in many plant tissues.

Silicon-burning process

In astrophysics, silicon burning is a very brief sequence of nuclear fusion reactions that occur in massive stars with a minimum of about 8-11 solar masses. Silicon burning is the final stage of fusion for massive stars that have run out of the fuels that power them for their long lives in the main sequence on the Hertzsprung-Russell diagram. It follows the previous stages of hydrogen, helium, carbon, neon and oxygen burning processes.

Silicon burning begins when gravitational contraction raises the star's core temperature to 2.7–3.5 billion Kelvin (GK). The exact temperature depends on mass. When a star has completed the silicon-burning phase, no further fusion is possible. The star catastrophically collapses and may explode in what is known as a Type II supernova.


A star is an astronomical object consisting of a luminous spheroid of plasma held together by its own gravity. The nearest star to Earth is the Sun. Many other stars are visible to the naked eye from Earth during the night, appearing as a multitude of fixed luminous points in the sky due to their immense distance from Earth. Historically, the most prominent stars were grouped into constellations and asterisms, the brightest of which gained proper names. Astronomers have assembled star catalogues that identify the known stars and provide standardized stellar designations. However, most of the estimated 300 sextillion (3×1023) stars in the Universe are invisible to the naked eye from Earth, including all stars outside our galaxy, the Milky Way.

For at least a portion of its life, a star shines due to thermonuclear fusion of hydrogen into helium in its core, releasing energy that traverses the star's interior and then radiates into outer space. Almost all naturally occurring elements heavier than helium are created by stellar nucleosynthesis during the star's lifetime, and for some stars by supernova nucleosynthesis when it explodes. Near the end of its life, a star can also contain degenerate matter. Astronomers can determine the mass, age, metallicity (chemical composition), and many other properties of a star by observing its motion through space, its luminosity, and spectrum respectively. The total mass of a star is the main factor that determines its evolution and eventual fate. Other characteristics of a star, including diameter and temperature, change over its life, while the star's environment affects its rotation and movement. A plot of the temperature of many stars against their luminosities produces a plot known as a Hertzsprung–Russell diagram (H–R diagram). Plotting a particular star on that diagram allows the age and evolutionary state of that star to be determined.

A star's life begins with the gravitational collapse of a gaseous nebula of material composed primarily of hydrogen, along with helium and trace amounts of heavier elements. When the stellar core is sufficiently dense, hydrogen becomes steadily converted into helium through nuclear fusion, releasing energy in the process. The remainder of the star's interior carries energy away from the core through a combination of radiative and convective heat transfer processes. The star's internal pressure prevents it from collapsing further under its own gravity. A star with mass greater than 0.4 times the Sun's will expand to become a red giant when the hydrogen fuel in its core is exhausted. In some cases, it will fuse heavier elements at the core or in shells around the core. As the star expands it throws a part of its mass, enriched with those heavier elements, into the interstellar environment, to be recycled later as new stars. Meanwhile, the core becomes a stellar remnant: a white dwarf, a neutron star, or, if it is sufficiently massive, a black hole.

Binary and multi-star systems consist of two or more stars that are gravitationally bound and generally move around each other in stable orbits. When two such stars have a relatively close orbit, their gravitational interaction can have a significant impact on their evolution. Stars can form part of a much larger gravitationally bound structure, such as a star cluster or a galaxy.

Stellar nucleosynthesis

Stellar nucleosynthesis is the theory explaining the creation (nucleosynthesis) of chemical elements by nuclear fusion reactions between atoms within stars. Stellar nucleosynthesis has occurred continuously since the original creation of hydrogen, helium and lithium during the Big Bang. It is a highly predictive theory that today yields excellent agreement between calculations based upon it and the observed abundances of the elements. It explains why the observed abundances of elements in the universe grow over time and why some elements and their isotopes are much more abundant than others. The theory was initially proposed by Fred Hoyle in 1946, who later refined it in 1954. Further advances were made, especially to nucleosynthesis by neutron capture of the elements heavier than iron, by Margaret Burbidge, Geoffrey Burbidge, William Alfred Fowler and Hoyle in their famous 1957 B2FH paper, which became one of the most heavily cited papers in astrophysics history.

Stars evolve because of changes in their composition (the abundance of their constituent elements) over their lifespans, first by burning hydrogen (main sequence star), then helium (red giant star), and progressively burning higher elements. However, this does not by itself significantly alter the abundances of elements in the universe as the elements are contained within the star. Later in its life, a low-mass star will slowly eject its atmosphere via stellar wind, forming a planetary nebula, while a higher–mass star will eject mass via a sudden catastrophic event called a supernova. The term supernova nucleosynthesis is used to describe the creation of elements during the evolution and explosion of a pre-supernova massive star (12–35 times the mass of the sun). Those massive stars are the most prolific source of new isotopes from carbon (Z = 6) to nickel (Z = 28).

The advanced sequence of burning fuels is driven by gravitational collapse and its associated heating, resulting in the subsequent burning of carbon, oxygen and silicon. However, most of the nucleosynthesis in the mass range A = 28–56 (from silicon to nickel) is actually caused by the upper layers of the star collapsing onto the core, creating a compressional shock wave rebounding outward. The shock front briefly raises temperatures by roughly 50%, thereby causing furious burning for about a second. This final burning in massive stars, called explosive nucleosynthesis or supernova nucleosynthesis, is the final epoch of stellar nucleosynthesis.

A stimulus to the development of the theory of nucleosynthesis was the discovery of variations in the abundances of elements found in the universe. The need for a physical description was already inspired by the relative abundances of isotopes of the chemical elements in the solar system. Those abundances, when plotted on a graph as a function of atomic number of the element, have a jagged sawtooth shape that varies by factors of tens of millions (see history of nucleosynthesis theory). This suggested a natural process that is not random. A second stimulus to understanding the processes of stellar nucleosynthesis occurred during the 20th century, when it was realized that the energy released from nuclear fusion reactions accounted for the longevity of the Sun as a source of heat and light.

Supernova nucleosynthesis

Supernova nucleosynthesis is a stellar evolution theory about the origin of the natural abundances of the chemical elements as created in supernovae. It explains how the nucleosynthesis of elements, from carbon to nickel, are made in the cores of massive stars. Knowledge of this process was a major stepping-stone in understanding how supernovae deposit these elements into the surrounding environment by enriching the interstellar medium and providing recycling materials for the birth of new stars. Supernova nucleosynthesis was first postulated by Fred Hoyle in 1954.

Type II supernova

A Type II supernova (plural: supernovae or supernovas) results from the rapid collapse and violent explosion of a massive star. A star must have at least 8 times, but no more than 40 to 50 times, the mass of the Sun (M☉) to undergo this type of explosion. Type II supernovae are distinguished from other types of supernovae by the presence of hydrogen in their spectra. They are usually observed in the spiral arms of galaxies and in H II regions, but not in elliptical galaxies.

Stars generate energy by the nuclear fusion of elements. Unlike the Sun, massive stars possess the mass needed to fuse elements that have an atomic mass greater than hydrogen and helium, albeit at increasingly higher temperatures and pressures, causing increasingly shorter stellar life spans. The degeneracy pressure of electrons and the energy generated by these fusion reactions are sufficient to counter the force of gravity and prevent the star from collapsing, maintaining stellar equilibrium. The star fuses increasingly higher mass elements, starting with hydrogen and then helium, progressing up through the periodic table until a core of iron and nickel is produced. Fusion of iron or nickel produces no net energy output, so no further fusion can take place, leaving the nickel–iron core inert. Due to the lack of energy output creating outward thermal pressure, the core contracts due to gravity until the overlying weight of the star can be supported largely by electron degeneracy pressure.

When the compacted mass of the inert core exceeds the Chandrasekhar limit of about 1.4 M☉, electron degeneracy is no longer sufficient to counter the gravitational compression. A cataclysmic implosion of the core takes place within seconds. Without the support of the now-imploded inner core, the outer core collapses inwards under gravity and reaches a velocity of up to 23% of the speed of light and the sudden compression increases the temperature of the inner core to up to 100 billion kelvins. Neutrons and neutrinos are formed via reversed beta-decay, releasing about 1046 joules (100 foe) in a ten-second burst. Also, the collapse of the inner core is halted by neutron degeneracy, causing the implosion to rebound and bounce outward. The energy of this expanding shock wave is sufficient to disrupt the overlying stellar material and accelerate it to escape velocity, forming a supernova explosion. The shock wave and extremely high temperature and pressure rapidly dissipate but are present for long enough to allow for a brief period during which the

production of elements heavier than iron occurs. Depending on initial size of the star, the remnants of the core form a neutron star or a black hole. Because of the underlying mechanism, the resulting supernova is also described as a core-collapse supernova.

There exist several categories of Type II supernova explosions, which are categorized based on the resulting light curve—a graph of luminosity versus time—following the explosion. Type II-L supernovae show a steady (linear) decline of the light curve following the explosion, whereas Type II-P display a period of slower decline (a plateau) in their light curve followed by a normal decay. Type Ib and Ic supernovae are a type of core-collapse supernova for a massive star that has shed its outer envelope of hydrogen and (for Type Ic) helium. As a result, they appear to be lacking in these elements.

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